Astronomy
&
Astrophysics
A&A 540, A10 (2012)
DOI: 10.1051/0004-6361/201117934
c ESO 2012
Herschel observations of a potential core-forming clump:
Perseus B1-E
S. I. Sadavoy1,2 , J. Di Francesco2,1 , Ph. André3 , S. Pezzuto4 , J.-P. Bernard5,6 , S. Bontemps3 , E. Bressert7,8 ,
˜ Lu’o’ng3 , N. Peretto3 ,
S. Chitsazzadeh1,2 , C. Fallscheer1,2 , M. Hennemann3 , T. Hill3 , P. Martin9 , F. Motte3 , Q. Nguyên
9
3
7,10
11,12
13
, and C. Wilson
M. Reid , N. Schneider , L. Testi , G. J. White
1
2
3
4
5
6
7
8
9
10
11
12
13
Department of Physics and Astronomy, University of Victoria, PO Box 355, STN CSC, Victoria BC, V8W 3P6, Canada
e-mail: [email protected]
National Research Council Canada, Herzberg Institute of Astrophysics, 5071 West Saanich Road, Victoria BC, V9E 2E7, Canada
Laboratoire AIM, CEA/DSM-CNRS-Université Paris Diderot, IRFU/Service d’Astrophysique, CEA Saclay, 91191 Gif-sur-Yvette,
France
IAPS – INAF, via del Fosso del Cavaliere, 100, 00133 Roma, Italy
CNRS, IRAP, 9 Av. colonel Roche, BP 44346, 31028 Toulouse Cedex 4, France
Université de Toulouse, UPS-OMP, IRAP, 31028 Toulouse Cedex 4, France
ESO, Karl Schwarzschild-Strasse 2, 87548 Garching bei Munchen, Germany
School of Physics, University of Exeter, Stocker Road, Exeter EX4 4QL, UK
Canadian Institute for Theoretical Astrophysics, University of Toronto, Toronto, ON, M5S 3H8, Canada
INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy
The Rutherford Appleton Laboratory, Chilton, Didcot OX11 0NL, UK
The Open University, Department of Physics and Astronomy, Milton Keynes MK7 6AA, UK
Department of Physics and Astronomy, McMaster University, Hamilton, ON, L8S 4M1, Canada
Received 23 August 2011 / Accepted 28 November 2011
ABSTRACT
We present continuum observations of the Perseus B1-E region from the Herschel Gould Belt Survey. These Herschel data reveal a
loose grouping of substructures at 160−500 μm not seen in previous submillimetre observations. We measure temperature and column
density from these data and select the nine densest and coolest substructures for follow-up spectral line observations with the Green
Bank Telescope. We find that the B1-E clump has a mass of ∼100 M and appears to be gravitationally bound. Furthermore, of the
nine substructures examined here, one substructure (B1-E2) appears to be itself bound. The substructures are typically less than a
Jeans length from their nearest neighbour and thus, may interact on a timescale of ∼1 Myr. We propose that B1-E may be forming a
first generation of dense cores, which could provide important constraints on the initial conditions of prestellar core formation. Our
results suggest that B1-E may be influenced by a strong, localized magnetic field, but further observations are still required.
Key words. stars: formation – dust, extinction – ISM: individual objects: Perseus B1-E
1. Introduction
Molecular clouds are highly structured regions of dust and gas.
They contain dense, small-scale, star-forming “cores” (0.1 pc)
that are usually clustered into larger-scale clumps and organized
along filaments (Williams et al. 2000; Ward-Thompson et al.
2007; Di Francesco et al. 2007; André et al. 2010). The cause of
this hierarchical structure, where the large-scale clouds (∼10 pc)
with moderate densities (∼102 cm−3 ) produce filaments, clumps,
and dense cores at higher densities (104 cm−3 ), is not well understood (Bergin & Tafalla 2007).
Molecular clouds form dense cores when diffuse gas is compressed. Current theories for core formation primarily focus on
two ideas: that cores form via (1) turbulent compression of diffuse gas (e.g., Larson 1981; Mac Low & Klessen 2004; Dib et al.
2009); or (2) the motion of neutral material between magnetic
field lines or ambipolar diffusion (e.g., Mestel & Spitzer 1956;
Mouschovias 1976; Kunz & Mouschovias 2009). Both mechanisms and gravity likely play important roles in star formation,
but it is unclear whether one mechanism would dominate in all
situations (i.e., clustered or isolated star formation and low-mass
or high-mass star formation). Furthermore, theoretical studies
also suggest that additional processes, such as radiation feedback
(i.e., Krumholz et al. 2010) or large-scale shocks associated with
converging flows (i.e., Heitsch et al. 2011), can influence molecular cloud structure and thus, core formation. Additionally, recent Herschel studies have shown that filaments are very prominent in star forming regions and likely have an important role in
the formation of dense substructures in molecular clouds (e.g.,
André et al. 2010; Arzoumanian et al. 2011).
Surveys of core populations in molecular clouds indicate that
cores are confined to regions of high column density. Johnstone
et al. (2004) and Kirk et al. (2006) found a relationship between
core occurrence and extinction for the Ophiuchus and Perseus
clouds, respectively, suggesting that there is a core formation
threshold at AV 5. Similarly, Lada et al. (2009) and André
et al. (2010) each compared two clouds with different degrees of
star formation activity and each found that the more active cloud
was composed of higher column density material (by a factor
of ∼10) than the quiescent cloud. These studies emphasize that
cores require a minimum column density (extinction) of material to condense from the bulk cloud (see also Heiderman et al.
2010).
Article published by EDP Sciences
A10, page 1 of 13
A&A 540, A10 (2012)
The precursors to cores are difficult to identify. Dense cores
are often influenced by processes such as nearby outflows and radiation feedback from a previous epoch of nearby star formation,
and observations of them cannot be used to constrain the dynamic properties of the initial core-forming material (Curtis &
Richer 2011). Without knowing the initial conditions and processes that cause cores to form from diffuse gas, we cannot accurately model their formation or evolution. Thus, identifying and
analyzing a core forming region without earlier episodes of star
formation would be exceedingly useful to constrain how molecular clouds form star-forming substructures.
In this paper, we use observations from the Herschel Gould
Belt Survey to explore a ∼0.1 deg2 clump roughly 0.7◦ east of
the B1 clump (Bachiller & Cernicharo 1986) within the Perseus
molecular cloud. This region, hereafter called B1-E, has high extinction (AV > 5) similar to the core formation threshold. Despite
this high extinction, previous submillimetre and infrared continuum observations suggest that B1-E contains neither dense cores
nor young stellar objects (e.g., Enoch et al. 2006; Kirk et al.
2006; Jørgensen et al. 2007; Evans et al. 2009). In contrast, our
Herschel observations show substructure that was not detected
by these other far-infrared and submillimetre continuum studies.
Furthermore, we use recent Green Bank Telescope (GBT) observations to quantify in part the kinematic motions of the densest
substructures seen in the Herschel data.
We propose that B1-E is forming a first generation of dense
cores in a pristine environment. In Sect. 2, we describe our
Herschel and GBT observations. In Sect. 3 we describe the properties we derive from our data. In Sect. 4 we discuss the implication of our results and compare our observations to theoretical
models. Finally, in Sect. 5 we summarize the paper.
2. Data
2.1. Herschel observations
The western half of Perseus, including B1-E, was observed
by the Photodetector Array Camera and Spectrometer (PACS;
Poglitsch et al. 2010) and the Spectral and Photometric Imaging
Receiver (SPIRE; Griffin et al. 2010) as part of the Herschel
Gould Belt Survey (André & Saraceno 2005; André et al. 2010).
The Herschel observations were taken in “parallel mode”, resulting in simultaneous coverage at 5 wavelength bands, with
the 70 μm and 160 μm PACS channels and the 250 μm, 350 μm,
and 500 μm SPIRE channels, over roughly 6 square degrees.
The 70 μm observations of B1-E, however, are less sensitive due
to the fast scan rate (60 /s) and low emission from cold material
(see Sect. 3.1). Thus, we will not include the 70 μm observations
in our discussion. For a full explanation of the observations, see
Pezzuto et al. (in prep.).
The PACS and SPIRE raw data were reduced using HIPE
version 5.0 and reduction scripts written by M. Sauvage (PACS)
and P. Panuzzo (SPIRE) that were modified from the standard
pipeline. For SPIRE reduction, we used updated calibration information (version 4). The final maps were created using the
scanamorphos routine developed by Roussel (2011). For consistency with other published maps, we set the pixel scale1 to the
default size from the HIPE mapmaking tools, which results in
pixels of sizes 6.4 , 6.0 , 10.0 , and 14.0 for 160 μm, 250 μm,
350 μm, and 500 μm, respectively. We also adopt beam sizes
of 13.4 , 18.1 , 25.2 , and 36.6 for the 160 μm, 250 μm,
1
By default, scanamorphos sets a much smaller pixel scale.
A10, page 2 of 13
350 μm, and 500 μm bands respectively2 (see Griffin et al. 2010;
Poglitsch et al. 2010). Assuming a distance to Perseus of 235 pc
(Hirota et al. 2008), Herschel can detect structure on scales
of ∼0.04 pc at 500 μm. Figure 1 shows a three-colour image
of western Perseus with labels for B1-E and other prominent
subregions.
Unlike previous ground-based submillimetre instruments,
Herschel can detect large-scale diffuse emission (e.g., Schneider
et al. 2010; Miville-Deschênes et al. 2010; Arzoumanian et al.
2011). Additionally, Herschel has excellent sensitivity to lowlevel flux. For example, Fig. 2 compares SCUBA 850 μm observations of B1-E (smoothed to 23 resolution; Di Francesco et al.
2008) with the new SPIRE 250 μm observations (at 18 resolution). The SPIRE 250 μm data show prominent substructures
not identified in the SCUBA 850 μm data, though several faint
features at 850 μm appear to agree with the brighter structures
in the 250 μm map. With the higher sensitivity of Herschel, we
are able to identify clearly structures in B1-E that were too faint,
i.e., <3σ, to be robust detections with SCUBA.
2.2. GBT observations
We selected the nine highest peaks in our Herschel-derived H2
column density map of B1-E (see Sect. 3.2) for complementary
follow-up observations with the new K-band Focal Plane Array
(KFPA) receiver on the GBT3 . Our targets, hereafter named
B1-E1 to B1-E9 according to decreasing peak column density,
were observed on 03 March 2011 with single-pointings. We
used the KFPA receiver with one beam and four spectral windows to observe each target simultaneously in NH3 (1, 1), NH3
(2, 2), CCS (21 −10 ), and HC5 N (9–8) line emission at 23.6945
GHz, 23.7226 GHz, 22.3440 GHz, and 23.9639 GHz, respectively. Similar to Rosolowsky et al. (2008), we made frequencyswitching observations with 9-level sampling and 12.5 MHz
bandwidth over 4096 spectral channels in each window to
achieve a high velocity resolution of vch ≈ 0.0386 km s−1 at
23.69 GHz. B1-E1 to B1-E8 were observed for ∼1120 s, each.
B1-E9 was observed for ∼840 s.
The GBT data were reduced using standard procedures in
GBTIDL4 for frequency-switched data. In brief, individual scans
from each spectral window were filtered for spikes or baseline
wiggles, folded, and then averaged. Baselines were obtained for
the averaged spectra by fitting 5th-order polynomials to line-free
channel ranges at the low- and high-frequency edges of each
band, and were then subtracted. To improve the detection levels,
the data were smoothed with a boxcar kernel equal to two channels in width. The reduced data were exported from GBTIDL
into standard FITS files using routines of AIPS++ developed by
Langston. Further analysis was conducted in MIRIAD5 and IDL
(Interactive Data Language). The 1σ noise levels were typically
∼0.01−0.02 K. At 23.69 GHz, the GBT beam is ∼33 and the
beam efficiency is ηb = 0.825.
2
Due to the fast scan rate, the 160 μm beam is slightly elongated
√ along
the scan direction. Thus, we adopt the geometric mean (i.e., r = ab) as
the beam size, where the elongated beam dimensions are from Poglitsch
et al. (2010).
3
The GBT is an 100 m telescope operated by the National Radio
Astronomy Observatory.
4
GBTIDL is a special IDL package specific to the GBT.
5
Multichannel Image Reconstruction, Interactive Analysis, and
Display (MIRIAD) software is developed by the Berkeley Illinois
Maryland Array (BIMA) group. See Sault et al. (1995) for more details.
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
Fig. 1. Three-colour image of the western half of the Perseus molecular cloud. Colour mosaic was generated using Herschel 160 μm, 250 μm,
and 350 μm observations. The white box denotes our boundary for B1-E for subsequent figures. The other prominent clumps in Perseus are also
labeled.
Fig. 2. Observations of Perseus B1-E from a) SCUBA 850 μm and b) SPIRE 250 μm maps. The SCUBA data came from the Extended SCUBA
Legacy Catalogue (see Di Francesco et al. 2008). Contours represent extinction levels of AV = 5, 7, 8 mag from the COMPLETE extinction map
(Ridge et al. 2006a) and the filled circles show the beam sizes of 23 for smoothed SCUBA data at 850 μm and 18 for SPIRE data at 250 μm.
A10, page 3 of 13
A&A 540, A10 (2012)
Table 1. GBT target information.
Source
B1-E1
B1-E2
B1-E3
B1-E4
B1-E5
B1-E6
B1-E7
B1-E8
B1-E9
α (J2000)
(h:m:s)
3:35:55.0
3:36:04.4
3:35:52.1
3:35:51.0
3:36:37.3
3:36:41.0
3:36:39.0
3:36:05.0
3:36:18.5
δ (J2000)
(◦ )
31:14:16
31:11:47
31:15:39
31:12:34
31:11:41
31:15:05
31:14:27
31:14:28
31:14:31
Peak N(H2 )a
(1021 cm−2 )
18.5
16.7
15.1
14.9
14.5
13.7
13.6
13.1
12.9
Detectionsb
NH3
NH3
NH3
NH3
NH3
NH3
NH3
NH3
NH3
(1, 1), CCS (21 −10 )
(1, 1), NH3 (2, 2), CCS (21 −10 ), HC5 N (9–8)
(1, 1), CCS (21 −10 )
(1, 1)
(1, 1)
(1, 1)
(1, 1), CCS (21 −10 )
(1, 1), CCS (21 −10 )
(1, 1)
Notes. (a) Peak Herschel-derived H2 column density towards the sources. See Sect. 3 for more details.
source.
Table 1 gives the positions, peak H2 column density, and the
detected line emission of these targets. NH3 (1, 1) line emission
was detected (>3σ) towards all nine targets and CCS (21 −10 )
line emission towards several. The NH3 (2, 2) and HC5 N (9–8)
lines were only detected toward B1-E2.
3. Results
3.1. SED fitting to Herschel data
We corrected the arbitrary zero-point flux offset in each Herschel
band using the method proposed in Bernard et al. (2010) that is
based on a comparison with the Planck HFI (DR2 version, see
Planck HFI Core Team et al. 2011) and IRAS data. In addition,
each map was convolved to the resolution of the 500 μm map
(36.6) and regridded to 14 pixels. The map intensities of the
160−500 μm bands were then fit by the modified black body
function,
Iν = κν Bν(T )Σ
(1)
where κν is the dust opacity, Bν is the black body function, T is
the dust temperature, and Σ is the gas mass column density. Note
that Σ = μmH N(H2 ), where μ is the mean molecular weight, mH
is the hydrogen mass, and N(H2 ) is the gas column density. For
consistency with other papers in the Gould Belt Survey (e.g.,
André et al. 2010),
κν = 0.1(ν/1000 GHz)β cm2 g−1 ,
(2)
where β is the dust emissivity index. The SED fits were made
using the IDL program mpfitfun by Markwardt. In brief, mpfitfun performs a least-squares comparison between the data and a
model function by adjusting the desired parameters until a best
fit is achieved.
Fitting β requires a plethora of data, particularly along the
Rayleigh-Jeans tail (e.g., λ > 300 μm at 10 K), to remove degeneracies in the model fits (Doty & Leung 1994; Shetty et al. 2009).
Since we have only two photometric bands (350 μm and 500 μm)
along the Rayleigh-Jeans tail, we assume β = 2, consistent with
the adopted β in other Herschel first-look studies (e.g., André
et al. 2010; Arzoumanian et al. 2011). There is some evidence
for β ≈ 2 in recent Planck studies (see Planck Collaboration et al.
2011a,b, and references therein). Adopting β = 2 here provides
good fits to our data (see below).
Figure 3 shows the dust temperatures and column densities
across B1-E resulting from the modified black body fits to each
pixel for the 160−500 μm bands, assuming flux uncertainties of
A10, page 4 of 13
(b)
Line emission detected towards each
15% based on calibration uncertainties (see Griffin et al. 2010;
Poglitsch et al. 2010) and β = 2. Both temperature and column density in B1-E are highly structured (see Figs. 3a and b),
where regions of higher column density are also associated with
slightly cooler temperatures and regions of lower column density have slightly warmer temperatures. Most of the cooler, high
column density material (6 × 1021 cm−2 ) is clumped into a central ring-like structure (i.e., slight depression is seen towards the
very centre).
Figures 3c and d show the number histograms of temperature
and H2 column density, respectively. These distributions are nonGaussian. The sample mean and standard deviation about the
mean for the temperature and column density are 14.1 K ± 0.8 K
and (6.3 ± 2.7) × 1021 cm−2 , respectively, where the standard
deviation is computed from,
1 (xi − x̄)2 ,
SE =
(3)
N−1
where x̄ is the population mean for the sample of size N.
These mean values agree well with previously estimated
quantities for this region. For example, Schnee et al. (2005)
measured a mean dust temperature of ∼14 K for the B1-E region using the ratio of IRAS 60 μm and 100 μm flux densities
to estimate dust temperature. Additionally, extinction data from
the COMPLETE survey (see contours in Figs. 2 and 3) suggest
a mean column density of (5.3 ± 1.5) × 1021 cm−2 , assuming
N(H2 )/AV = 1021 cm−2 mag−1 . Using a higher resolution extinction map from S . Bontemps, we find a mean column density
of (6.3±3.6)×1021 cm−2 . These two column densities agree very
well with our measured value of (6.3 ± 2.7) × 1021 cm−2 and they
are also similar to the threshold column density for dense core
formation from the extinction analysis in Perseus by Kirk et al.
(2006), 5 × 1021 cm−2 . For comparison, Fig. 3d includes the Kirk
et al. column density threshold as a dashed line. Although B1-E
does not appear filamentary, there is ample material from which
dense cores may form. With higher resolution data, André et al.
(2011) found a threshold column density of 7 × 1021 cm−2 for
dense structures to form in Aquila via thermal instabilities along
a filament at 10 K.
3.2. Column density profiles
One of the more prominent core models is the Bonnor-Ebert
sphere (Bonnor 1956; Ebert 1955), which represents the density profile of a sphere in hydrostatic equilibrium under the influence of an external pressure. The Bonnor-Ebert sphere has
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
Fig. 3. SED-fitting results for B1-E. Top panels shows the temperature a) and H2 column density b) maps across B1-E as measured from SEDfitting to 160–500 μm data, according to Eqs. (1) and (2) and assuming β = 2. Bottom panels show the histograms of the above maps. The
temperature histogram c) uses a bin size of 0.3 K and the H2 column density histogram d) uses a bin size of 1 × 1021 cm−2 . For comparison, the
dashed line in panel d) indicates the observed column density threshold (from extinction) for core formation from Kirk et al. (2006).
a flat inner density distribution and a power-law density downtrend at larger radii. Many prestellar cores, i.e., dense cores that
are gravitationally bound but do not show evidence of a central
luminous protostar, have shown Bonnor-Ebert-like profiles (e.g.,
Ward-Thompson et al. 1999; Alves et al. 2001).
With the excellent spatial resolution of Herschel, we can
measure the column density profiles for individual B1-E
substructures. First, we identified the locations of peak column density using the 2D Clumpfind algorithm (Williams et al.
1994). Briefly, Clumpfind identifies intensity peaks and then
uses closed contours at lower intensity levels to assign boundaries. We will discuss the nine highest column density substructures identified by Clumpfind (B1-E1 to B1-E9) for this paper.
Second, we measured the azimuthally-averaged column density
profile of our nine sources using the ellint task in MIRIAD. For
simplicity, we used circular annuli of 10 width for r > 7 . For
the central radii (r < 7 ), we assume the peak column density.
We caution that several of our sources appear elliptical and our
circular approximation is meant to provide a broad, first-look
analysis.
Figure 4 shows the column density profiles, where the column density values are plotted from the centre of each annulus.
For comparison, Fig. 4 also shows the beam profile (solid grey
curve), a “generic” column density profile convolved with the
beam (dashed grey curve) and an analytical Bonnor-Ebert profile
convolved with the beam (dotted curve). For both analytic profiles, we follow the approximation from Dapp & Basu (2009)
and assume a temperature of T = 10 K, a central density of
n = 106 cm−3 and a constant of proportionality of k = 0.54,
where k ≈ 0.4 for a singular isothermal sphere and k ≈ 1
for a collapsing cloud. Our observed column density profiles
are much wider than the models, likely due to contaminating
material in the foreground or far-background along the line
of sight, hereafter called LOS material. Towards the centre of
each source, the analytic models and observed profiles are more
similar, since the source column density dominates over the
LOS level, whereas in the power-law roll-off, the LOS material is
likely more significant and the profiles deviate from the analytic
models. Unfortunately, the column density of such extended
material is difficult to distinguish from the source column density. To estimate its contribution towards each source, we used
the subsequent column density at the location where the powerlaw slope in the column density profile flattened or began to increase. Thus, the LOS material level ranges from ∼8 ×1021 cm−2
for B1-E9 to ∼11 × 1021 cm−2 for B1-E1 or roughly 60% of the
peak column density. These LOS column densities are very conservative and could overestimate their actual contributions by as
much as a factor of 2. Indeed, these values are larger than the
mean B1-E column density of 6.3 × 1021 cm−2 .
Figure 5 shows a comparison of our nine column density
profiles after correcting each profile for LOS material and normalizing to the respective peak column density, and includes
the Bonnor-Ebert profile, generic profile, and beam profile from
Fig. 4. All nine profiles follow a shape similar to those of the analytic models, with a flat centre and steep falloff towards larger
radii. Additionally, all profiles but B1-E5 generally appear more
centrally concentrated than the models, implying that B1-E5 is
the least compact. The column density profiles do show some
differences with the models, however, such as steeper fall-offs.
3.3. Substructures
We estimated source sizes and masses based on Gaussian fits
to the LOS-subtracted column density profiles. Source size was
defined by the FWHM of the Gaussian (i.e., R = FWHM/2).
Since B1-E6 and B1-E7 have a small projected separation
A10, page 5 of 13
A&A 540, A10 (2012)
Fig. 4. The azimuthally-averaged column density profiles for the nine sources in B1-E, assuming a distance of 235 pc. The profiles were measured
using circular annuli with a thickness of 10 for r > 7 . Note that the central area is defined by a circle of radius 7 but plotted at 3.5 . The
dashed line indicates our estimate for the non-source line-of-sight (LOS) material. For comparison, the dotted curve and dashed curve illustrate
a Bonnor-Ebert column density profile and a “generic” [1 + (r/a)2 ]−0.5 column density profile following Dapp & Basu (2009). For both analytic
curves, we assume a temperature of 10 K and a central density of n = 106 cm−3 and convolved the profiles with a 36.6 beam. The grey solid curve
shows the beam profile. The analytical profiles and the beam profile are scaled to the peak column density. Note that both axes are logarithmic.
of ∼50 , we truncated both to 31 radii6 to ensure that we mostly
measured column densities unique to each. Thus, our analyses of
B1-E6 and B1-E7 are biased toward the denser, central regions
unlike those of the other sources. Subtracting out the LOS material limits the influence from nearby extended material and generally fits only the denser, embedded object. If we include the
LOS material, the estimated source sizes generally increase by
20%. Isolated substructures, like B1-E2 or B1-E8, have sizes
that vary by ∼5% and substructures embedded in extended profiles like, B1-E1 or B1-E4, have sizes that vary by 24% and 37%,
respectively.
Table 2 lists the adopted properties for our sources, including their deconvolved radii, estimated (upper and lower) masses,
and average densities, assuming they are perfect spheres and dividing the (upper and lower) source masses with their spherical
volumes, 43 πR3d , where Rd is the deconvolved radius. We consider two mass and density limits: the lower limits subtract out
the LOS material and the upper limits include the LOS material. Note that, by our definition of source size, we are measuring the inner regions of each profile where core precursors are
most likely to arise. Since we are using the Gaussian FWHM
to estimate the source size, the true source sizes, including any
diffuse envelopes, may be larger by a factor of 2. Source mass
was estimated by summing over the column density and assuming a mean molecular weight of μ = 2.33. These results provide
6
Deconvolved radii are ≈25 .
A10, page 6 of 13
a reasonable first look at the relative sizes and masses of these
sources.
3.4. Line emission
Kirk et al. (2007) and Rosolowsky et al. (2008) each attempted
to identify dense gas towards B1-E via high spatial resolution
line observations in N2 H+ and NH3 , respectively. They observed
several “blind” pointings towards B1-E and found no strong detections. Unlike previous continuum studies, our Herschel data
have identified several cold, dense substructures in B1-E (see
Fig. 3) and these were not directly probed by either Kirk et al.
or Rosolowsky et al. Using our Herschel results, we selected
the nine substructures with the highest column densities for
follow-up observations with the GBT in NH3 (1, 1), NH3 (2, 2),
CCS (21 −10 ), and HC5 N (9–8) line emission. Figure 6 shows the
locations of our nine targets (see also Table 1).
3.4.1. NH3
We detected NH3 (1, 1) emission towards all nine targets.
We fit the NH3 (1, 1) hyperfine lines with multiple Gaussian
components using the following equation;
τ(v) = τ1,1
18
i=1
v − v − v 2 i
lsr
αi exp −4 ln 2
Δv
(4)
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
Fig. 5. Normalized column density profiles for sources in B1-E, assuming a distance of 235 pc. Column density was determined from azimuthallyaveraged annuli (see Fig. 4). Furthermore, the LOS material was estimated for each source and subtracted from the azimuthal average. The grey
solid, dotted, and dashed curves are the same as for Fig. 4.
Table 2. Column density determined properties.
Source
B1-E1
B1-E2
B1-E3
B1-E4
B1-E5
B1-E6d
B1-E7d
B1-E8
B1-E9
Rd a
(AU)
7.1 × 103
6.9 × 103
7.7 × 103
6.0 × 103
9.3 × 103
5.9 ×103
5.9 ×103
7.2 × 103
5.4 × 103
Mb
(M )
0.6−1.6
0.6−1.4
0.5−1.5
0.3−1.0
0.5−2.0
0.3−1.0
0.3−0.9
0.3−1.2
0.3−0.8
nc
(104 cm−3 )
5.9−17
6.7−15
4.0−12
4.3−18
2.4−9.1
5.7−17
5.6−17
3.1−12
5.8−18
Notes. (a) Object radii were estimated from Gaussian fits to the column density profile (R = FWHM/2). Radii were deconvolved with a
36.6 beam. (b) Masses are taken from integrating out the column density profiles. The lower limit measurements subtract the LOS material
from the column density profiles. The upper limit measurements do not
subtract the LOS material from the column density profiles. (c) Average
density determined from the mass and volume. (d) B1-E6 and B1-E7 are
truncated to measure column density unique to each source.
where τ1,1 is the optical depth for the (1, 1) transition, αi is the
transition weight, vi is the velocity of the hyperfine line (for
a given rest frequency), vlsr is the velocity of the source with
respect to the local standard of rest, and Δv is the FWHM of
the line. We used the hyperfine frequencies and weights from
Kukolich (1967) and Rydbeck et al. (1977) and derived the values for τ1,1 , vlsr , and Δv from the fits. Figure 7 shows example fits to the NH3 spectra. Table 3 lists centroid velocity (vlsr )
and uncorrected velocity line width (Δv) from the NH3 (1, 1)
fits, the kinetic gas temperature from NH3 (1, 1) and (2, 2) line
Fig. 6. Locations of substructures in B1-E overlaid on our Herschelderived column density map. Numbers indicate the positions of peak
column densities and the relative magnitude of the column density peak
(see also Table 1). The thin black circles show the deprojected source
sizes (see Sect. 3.3). The dark filled circle shows the GBT beam (∼33 )
at 23.69 GHz.
comparisons, and the resulting thermal and non-thermal velocity
dispersions for all nine sources.
For B1-E2, we could derive the kinetic gas temperature using NH3 (1, 1) and NH3 (2, 2) line emission, whereas we could
only determine upper limits for our other sources. We use our
NH3 -derived kinetic temperature of T K = 10.26 K ± 0.36 K for
B1-E2 and adopt the mean NH3 -derived kinetic temperature of
11 K from the survey of Rosolowsky et al. (2008) for our other
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A&A 540, A10 (2012)
Fig. 7. Fits to NH3 (1, 1) and (2, 2) spectra. Note the hyperfine components visible for the NH3 (1, 1) spectrum of B1-E2. The NH3 (2, 2) spectra
were smoothed by a boxcar with a 2-channel width to improve the visual appearance. For comparison, we show two undetected NH3 (2, 2)
observations. (This figure is available in color in the electronic form.)
Table 3. Properties from NH3 (1, 1) and NH3 (2, 2) line spectra.
Source
B1-E1
B1-E2
B1-E3
B1-E4
B1-E5
B1-E6
B1-E7
B1-E8
B1-E9
vlsr
(km s−1 )
7.61 ± 0.02
7.67 ± 0.002
7.47 ± 0.01
7.83 ± 0.01
7.57 ± 0.01
7.61 ± 0.01
7.31 ± 0.01
6.78 ± 0.01
7.11 ± 0.01
Δva
(km s−1 )
0.998 ± 0.056
0.293 ± 0.005
0.681 ± 0.030
1.078 ± 0.036
0.833 ± 0.022
0.905 ± 0.032
0.708 ± 0.032
0.537 ± 0.021
0.690 ± 0.038
TKb
(K)
<27
10.26 ± 0.36
<16
<27
<19
<23
<24
<20
<32
σT c
(km s−1 )
0.073 ± 0.007
0.071 ± 0.001
0.073 ± 0.007
0.073 ± 0.007
0.073 ± 0.007
0.073 ± 0.007
0.073 ± 0.007
0.073 ± 0.007
0.073 ± 0.007
σNT
(km s−1 )
0.417 ± 0.025
0.101 ± 0.003
0.280 ± 0.015
0.451 ± 0.016
0.344 ± 0.012
0.376 ± 0.015
0.290 ± 0.016
0.213 ± 0.012
0.282 ± 0.018
Notes. (a) The velocity FWHM from the best-fit models uncorrected for channel resolution. (b) Kinetic temperature towards each source. For B1-E2,
we measured the kinetic temperature from NH3 (1, 1) and NH3 (2, 2) line emission. For our other sources, we used a 3σ upper limit for NH3 (2, 2)
to estimate the upper limit of T K . (c) The thermal line width of B1-E2 was determined using the derived kinetic temperature, T K = 10.26 K ± 0.36 K.
For the remaining sources, we adopted 11 K, the mean NH3 -derived kinetic temperature from Rosolowsky et al. (2008) with an uncertainty of 2 K
from the dispersion of the kinetic temperature in that survey (see Enoch et al. 2008).
A10, page 8 of 13
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
sources. As an estimate of the uncertainty, we use the dispersion
in the kinetic temperatures from the Rosolowsky et al. data, 2 K,
(see Enoch et al. 2008). Note that these gas temperatures correspond to the denser, colder inner layers where ammonia is excited, whereas the dust temperatures (see Sect. 3.1) represent an
average of all the material along the line of sight, i.e., including
the lower density envelopes. For more information on extracting
kinetic temperatures from NH3 line emission, see Friesen et al.
(2009).
Table 3 also lists the thermal velocity dispersion, σT , and the
non-thermal velocity dispersion, σNT . We determined σT from
the kinetic temperature and σNT from the line widths via,
kb T K
(5)
σT =
μNH3 mH
σNT = σ2obs − σT 2 ,
(6)
where kb is the Boltzmann constant, T K is the kinetic temperature, μNH3 = 17.03 is the mean molecular weight of NH3 , mH
is the atomic hydrogen mass, and σobs is the√total corrected velocity dispersion of the lines, σobs = Δvc / 8 ln 2 for Δvc 2 =
Δv2 − v2ch . Using the NH3 study by Rosolowsky et al. (2008), we
adopted T K = 11 K for all our targets except B1-E2, where we
used our derived gas temperature, T K = 10.26 K ± 0.36 K.
3.4.2. CCS and HC5 N
CCS (21 −10 ) was detected towards five targets (see Table 1) and
HC5 N (9–8) was detected towards only B1-E2. Figure 8 shows
examples of these spectra. We fit the CCS (21 −10 ) and HC5 N
(9–8) spectra with single Gaussians. The velocities determined
for most CCS (21 −10 ) detections and the single HC5 N (9–8) detection were similar to those found from NH3 , suggesting these
dense gas tracers are probing the same region as the NH3 emission. The CCS (21 −10 ) spectra for B1-E1 and B1-E8, however,
revealed fairly different values. For B1-E1, the CCS (21 −10 )
line emission appears redshifted by ∼0.3 km s−1 (15σ) with respect to the NH3 (1, 1) line emission, whereas for B1-E8, the
CCS (21 −10 ) line emission is redshifted by ∼0.11 km s−1 (2σ).
The redshifted CCS (21 −10 ) emission in B1-E1 may be tracing
a different region of dense gas from the NH3 emission, such as a
region undergoing infall (e.g., see Swift et al. 2005).
4. Discussion
Our Herschel data have revealed extensive substructures within
B1-E for the first time. Such faint substructures may represent
an early evolutionary stage in star formation, where a first generation of cores is forming from lower density material. Thus,
it is important to characterize the substructures seen in these
Herschel maps.
4.1. Comparison with Jeans instability
The minimum length scale for gravitational fragmentation in a
purely thermal clump or cloud is described by the Jeans length,
λJ = (c2s π/Gρ)1/2 , where cs is the thermal sound speed, G is the
gravitational constant, and ρ is the density of material (Stahler
& Palla 2005). In terms of temperature and number density, we
find,
T 0.5 −0.5
n
λJ = 0.21
pc.
(7)
−3
10 K
104 cm
Ideally, the Jeans length should be calculated for the gas temperature and gas density. In small, dense scales (i.e., 105 cm−3 ;
Di Francesco et al. 2007), the gas and dust temperatures are coupled, whereas in larger, more diffuse scales, the gas and dust
are decoupled and the gas may be warmer than the dust (Young
et al. 2004; Di Francesco et al. 2007; Ceccarelli et al. 2007).
The central region of B1-E containing the substructures, i.e., the
−2
area contained within N(H2 ) > 6 × 1021 cm
√ , has a mass of
∼110 M and an effective radius (Reff = A/π) of ∼0.46 pc.
Approximating B1-E as a sphere, we find an average density of
∼5 × 103 cm−3 , suggesting that the the gas and dust in B1-E are
likely decoupled such that the mean dust temperature of ∼14 K
(see Sect. 3.1) is a lower limit to the gas temperature. Assuming
a lower limit gas temperature of 14 K and n = 5 × 103 cm−3 ,
the lower limit λJ ≈ 0.35 pc. Note that the kinetic temperatures
measured in Sect. 3.4.1 correspond to the smaller, denser scales
where the gas and dust temperatures should be coupled and not
the scales measured here.
The minimum projected separations between our nine substructures are generally between 0.1 pc and 0.19 pc, with a median value of ∼0.13 pc for all nine sources. Recall, however, that
B1-E6 and B1-E7 have an angular separation of ∼0.05 pc, a factor of 2 closer than any other pair of substructures. Despite their
close proximity, B1-E6 and B1-E7 show different centroid velocities (see Table 3), and therefore, cannot be considered a single object. Also, we can only measure the projected 2-D separations between our nine substructures and not their physical 3-D
separations. Assuming a typical inclination angle of 60◦ , our median separation is ∼0.15 pc for all nine substructures or ∼0.17 pc
excluding B1-E6 and B1-E7.
The median minimum separation is a factor of two smaller
than our best estimate of the Jeans length, though our Jeans
length estimate is, itself, uncertain within a factor of a few.
The fact that we find a minimum separation that is less
than the Jeans length is still interesting and could indicate
gravitational fragmentation followed by bulk contraction of the
B1-E group, but the overall uncertainties in determining the minimum separations and Jeans length make this difficult to validate. Additionally, B1-E does not appear to be filamentary (see
Fig. 6) and therefore, the clump could be extended along the line
of sight resulting in more significant radial distances between
substructures.
For the substructures, we use the Jeans mass to explore their
thermal stabilities. Similar to the Jeans length, the Jeans mass is
a critical mass scale for gravitational fragmentation (Stahler &
Palla 2005). This critical mass can be described as,
MJ = 2.9
T 1.5 −0.5
n
M .
10 K
104 cm−3
(8)
The B1-E substructures have average densities of n ∼ 1 ×
105 cm−3 (see Sect. 3.3). Furthermore, we found a gas temperature for B1-E2 of ∼10 K. Note, that this gas temperature corresponds to the denser B1-E2 substructure and should not be
assumed for the lower density B1-E clump. Assuming T = 10 K
and n = 1 × 105 cm−3 , the critical Jeans mass is MJ ≈ 0.9 M .
Several of the substructures have mass limits on order of the
Jeans mass (see Table 2), suggesting that these sources are approaching a critical, unstable mass. Thus, the B1-E substructures
are interesting prospects for future studies of substructures evolution.
A10, page 9 of 13
A&A 540, A10 (2012)
Fig. 8. Fits to the detected CCS (21 −10 ) and HC5 N (9–8) spectra. The weaker CCS (21 −10 ) spectra and the B1-E2 HC5 N (9–8) spectra were all
smoothed by a boxcar with a 2-channel width to improve the visual appearance of the lines. The B1-E2 CCS (21 −10 ) spectrum was not smoothed.
See the electronic edition of the Journal for a colour version of this figure.
4.2. Time scale for interactions
From the Herschel observations, the B1-E substructures appear
linked (see Fig. 6), which might indicate that they are interacting.
The timescale for any two substructures to interact can be estimated from the typical separation between the substructures and
their velocity dispersion. For simplicity, we adopt the average
projected separation between all nine substructures of ∼0.40 pc.
For an inclination angle of 60◦ , the typical separation between
the substructures is dave ∼ 0.46 pc. To estimate the velocity dispersion of the sources, we use the dispersion of the centroid velocities of NH3 (1, 1) from all nine objects. The weighted 1D
centroid velocity dispersion is σ1D ≈ 0.24 km s−1 and assuming symmetry, the 3D velocity dispersion7 is σ3D ≈ 0.42 km s−1 .
Thus, the B1-E substructures have an interaction timescale of
∼1 Myr, a factor of two larger than recent estimated prestellar core lifetimes and less than the expected protostellar core
lifetime (e.g., Kirk et al. 2005; Ward-Thompson et al. 2007;
Enoch et al. 2008). Furthermore, this timescale corresponds to
the interaction time between any two substructures and not the
interaction time between nearest neighbour substructures. Given
that the nearest neighbour separation is a factor of ∼3 smaller
than the typical separation adopted here, these substructures
7
Assuming σ23D = σ2r + σ2φ + σ2θ = 3σ2r .
A10, page 10 of 13
have the potential to interact prior to forming stars. Thus, competitive accretion may be significant during further evolution of
B1-E (e.g., Bonnell et al. 2001; Krumholz et al. 2005; Bonnell
& Bate 2006).
4.3. Comparison with virial equlibrium
Not all dense substructures will form cores and stars. For example, a core with kinetic energy from internal motions that exceeds its gravitational binding energy may be transient and unable to form persistent dense structures or stars. According to
McKee (1999), a core is gravitationally bound if the virial parameter, α, is
Mvirial 5σ2v R
=
2
(9)
M
GM
where σv is the velocity dispersion, R is the radius, and M is the
mass.
B1-E has an effective radius of 0.46 pc (see Sect. 4.1).
We estimate the velocity dispersion using the centroid velocity
dispersion from Sect. 4.2, σ3D = 0.42 km s−1 . This velocity
dispersion was derived from the motions of the individual substructures and thus, reflects the turbulence in the cloud that first
created the substructures rather than the thermal velocity dispersion of the gas. The expected thermal velocity of the gas is
α=
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
Table 4. Virial analysis properties.
Source
B1-E1
B1-E2
B1-E3
B1-E4
B1-E5
B1-E6
B1-E7
B1-E8
B1-E9
σv
(km s−1 )
0.461 ± 0.030
0.216 ± 0.005
0.342 ± 0.022
0.492 ± 0.022
0.397 ± 0.018
0.424 ± 0.022
0.351 ± 0.023
0.291 ± 0.021
0.344 ± 0.025
σv /cs a
2.3 ± 0.3
1.1 ± 0.03
1.7 ± 0.2
2.5 ± 0.3
2.0 ± 0.2
2.1 ± 0.2
1.8 ± 0.2
1.5 ± 0.2
1.7 ± 0.2
Mvirial
(M )
8.5 ± 4.8
1.8 ± 0.9
5.1 ± 2.9
8.2 ± 4.5
8.3 ± 4.5
6.0 ± 3.3
4.1 ± 2.3
3.4 ± 2.0
3.6 ± 2.1
αb
5–14
1–3
3–10
8–27
4–17
6–20
5–13
3–11
5–12
Notes. (a) Velocity dispersion divided by sound speed, assuming T =
11 K for all substructures except B1-E2 where T = 10.23 K. (b) Virial
parameter considering the upper and lower mass limits for each substructure (see Table 2).
less than 0.42 km s−1 , and thus non-thermal support is important. Assuming R = 0.46 pc and σv = 0.42 km s−1 , B1-E has a
virial mass of ∼96 M , or α ≈ 0.9 for a clump mass of 110 M
(see Sect. 4.1). These results imply that B1-E is itself gravitationally bound and thus, likely to form dense cores and stars in
the future.
For the individual substructures, velocity dispersions were
determined from NH3 (1, 1) line widths, after estimating the
non-thermal component (σNT ; see Table 3) and the thermal component (cs ) for a mean molecular weight of μH2 = 2.33. Table 4
shows our values for the total velocity dispersion, σv , the ratio of the velocity dispersion to the thermal sound speed, the
virial mass, Mvirial , and the virial parameter, α, for each substructure. Most of the virial parameters are 2, suggesting that these
substructures themselves are not gravitationally bound and thus,
may not necessarily represent persistent objects. B1-E2 has the
smallest virial parameter, 1 <
∼ α <
∼ 3, however, and may be itself gravitationally bound. B1-E8 and B1-E3 have lower limit
estimates of α ∼ 3 and given that we know α only within a factor of a few, these sources may also be gravitationally bound.
These virial limits, however, depend greatly on the source mass,
which can vary by a factor of ∼4 depending on our treatment
of the LOS material (see Sect. 3.3). To be bound, these candidates must have negligible LOS material towards them, which is
very unlikely. Thus, B1-E2 is the only strong candidate for being
gravitationally bound.
Most B1-E substructures have very large virial parameters
and are thus, expected to be unbound. In comparison, the Jeans
mass (see Sect. 4.1) is a mass scale for only gravitational fragmentation in a thermal clump and does not include the contributions from non-thermal support. Several of the substructures
with α 2 have masses on order of the Jeans mass, ∼0.9 M ,
suggesting that the Jeans mass analysis is too simplistic and does
not well represent the evolutionary state of these structures. For
example, the substructures may have significant turbulent motions. With only one potentially bound substructure (B1-E2), it is
difficult to determine if B1-E will eventually form several cores
or no cores at all (i.e., these substructures are all transient).
The substructures themselves show large line widths (see
Table 4), atypical of dense prestellar cores, which are subsonic.
In particular, B1-E2 has the narrowest line profiles (by at least
a factor of 2) and the strongest line detections, suggesting that
this object is the most evolved of the substructures. These observations hint at a possible dynamic evolution, where dense
cores first contain significant non-thermal, turbulent motions
which dissipate into coherent, quiescent dense cores. If so, the
substructures in B1-E may become more bound over time as
their turbulent support is dissipated, a process which may explain the narrower lines toward B1-E2.
Note that we assume T K = 11 K for our substructures (except B1-E2). In Sect. 3.4.1, we found upper limits for kinetic
temperature >11 K (see Table 3). Furthermore, we are using a
simplified virial equation in Eq. (9). Ideally, the virial equation
includes magnetic pressure and surface pressure terms which can
either assist gravity in the collapse or help in the support of a
clump (e.g., see Nakamura & Li 2008). Our observations, however, do not measure these quantities and thus, we use a simplified virial equation. Nevertheless, B1-E2 is special with respect
to the other substructures. Indeed, we may have the first observations of a core precursor.
4.4. Comparison with other star forming regions
B1-E has column densities well above ∼5 × 1021 cm−2 , the core
formation threshold reported by Kirk et al. (2006). Since the
other clumps with similarly high column densities are actively
forming stars, it is reasonable to expect B1-E to do the same in
the near future. Furthermore, Lada et al. (2010) found a good
correlation between the number of YSOs in a cloud (or clump)
with the cloud mass above AK > 0.8 (or AV > 7.3) for several
star forming regions. Based on their extinction threshold, we find
a clump mass of ∼90 M for N(H2 ) > 7.3 × 1021 cm−2 and thus,
we could expect ∼10 YSOs to form in B1-E. With Herschel,
we have detected several substructures with moderate densities
of ∼105 cm−3 (see Table 2), similar to what is often defined for
4
−3
dense cores, i.e., >
∼10 cm (Bergin & Tafalla 2007).
These substructures, however, were not well observed by
SCUBA at 850 μm or Bolocam at 1 mm (see Fig. 2), and thus,
are not dense cores in the traditional definition. The substructures in B1-E also have supersonic velocity dispersions. In general, low-mass starless cores have subsonic velocity dispersions
(e.g., Myers 1983; Kirk et al. 2007; André et al. 2007; Pineda
et al. 2010) whereas protostellar cores can have supersonic velocity dispersions (e.g., Gregersen et al. 1997; Di Francesco et al.
2001). Since we see no evidence of protostellar activity whatsoever (i.e., from Spitzer data), these substructures are likely unevolved.
Thus, the substructures in B1-E are likely core precursors
and more observations of this region should provide some insight into core formation (transient or persistent) in molecular
clouds. In particular, B1-E is isolated from YSOs and young
stars, so the region is unaffected by environmental processes
such as outflows or winds. A nearby expanding dust shell, seen
in IRAS 60 μm and 100 μm data (Ridge et al. 2006b), may be
sweeping material towards B1-E, though the ring itself does not
appear to be directly interacting with this clump. This shell may
have a larger impact on the evolution of structure in B1-E in the
future.
A relatively pristine core forming region is exceedingly rare.
The only additional case may be the inactive high extinction
clump (211) identified in L1495 by Schmalzl et al. (2010),
though further detailed studies of that clump are still necessary to
determine what, if any, substructures are forming cores and stars.
A collection of starless cores in Aquila form a similar grouping
as in B1-E, however this region has two known protostars nearby
(1 pc) and so is not as pristine as B1-E (Könyves et al. 2010,
see their Fig. 5).
4.5. Comparison with core formation models
Core formation simulations attempt to recreate molecular
cloud fragmentation, implementing factors such as gravitational
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A&A 540, A10 (2012)
collapse, magnetic pressure, and turbulent energy. Many recent
studies have argued that models of core formation must consider
both ambipolar diffusion and turbulence (e.g., Li & Nakamura
2004; Dib et al. 2007; Nakamura & Li 2008; Basu et al. 2009a,b;
Gong & Ostriker 2011). Magnetic fields introduce an additional
pressure support that opposes gravitational collapse, whereas
turbulence either opposes collapse by introducing kinetic energy
support or induces collapse via energy dissipation from compressive shocks.
Several recent studies have examined the effects of different magnetic field and turbulent flow properties (e.g., Li &
Nakamura 2004; Basu et al. 2009a; Price & Bate 2009). In particular, these simulations examine fragmentation in subcritical
(where magnetic fields dominate over gravity) and supercritical
(where gravity dominates over magnetic fields) regimes with dissipating turbulent flows. In the subcritical simulations, the fragmenting clouds generally form cores in relative isolation and
with less filamentary morphologies, such as ring-like distributions of loose core groups (see also Li & Nakamura 2002). These
cores tend to form more slowly, i.e., fragmentation is suppressed
by the strong magnetic field, and since large-scale turbulent
flows are generally damped, the velocity field is generally subsonic. (Note that Basu et al. 2009a, find that strong magnetic
fields can induce oscillations and create a supersonic velocity
field, i.e., 3cs .) In contrast, supercritical simulations generally
form cores associated with filaments, and these cores tend to
form relatively quickly and can move supersonically through
their environment.
An isolated core forming region is necessary to test model
predictions and constrain initial conditions of dense gas before the onset and influence of nearby stars and internal protostars. Thus, B1-E is an ideal target for such measurements (see
Sect. 4.4). From our observations, B1-E contains a small, loose
grouping of substructures. Although the substructure morphology observed in B1-E is not entirely ring-shaped, the distribution hints at an inclined ring-like structure. Furthermore, if star
formation is triggered, it is surprising that B1-E, which is bookended by two highly active and more evolved clumps, IC 348
to the far east and NGC 1333 to the far west, contains no evidence for evolved star formation itself. Furthermore, IC 348
shows many evolved pre-main sequence stars and NGC 1333,
though appearing younger, has many YSOs (Bally et al. 2008;
Herbst 2008; Walawender et al. 2008), suggesting that B1-E is
not part of an age gradient.
These observations suggest that B1-E is influenced by a
strong, localized magnetic field. Further tests of this scenario
would be possible by observing the entire velocity field for B1-E.
Instead, we could only estimate the bulk velocity field from the
relative motions of the individual substructures (from line centroids; see Sect. 4.2), σ3D = 0.42 km s−1 . Additionally, we do
not have direct measurements of the gas temperature for the
B1-E clump. If we use the mean B1-E dust temperature of 14 K
as the lower limit gas temperature, cs 0.22 km s−1 and the
velocity field within B1-E may be supersonic by <
∼2cs . These
motions can arise in either subcritical or supercritical clouds,
though the latter is generally associated with supersonic velocities.
Direct measurements of the magnetic field would also be
useful to constrain its role in B1-E. Unfortunately, very little
magnetic field information is available here. Goodman et al.
(1990) measured optical polarization towards all of Perseus, and
for B1-E, they found strong, ordered polarization vectors reaching ∼9% polarized light, a factor of two higher than any other
dense region in Perseus. While these polarization observations
A10, page 12 of 13
may suggest a strong magnetic field is associated with the
B1-E region, these observations are dependent on the field viewing angle. Also, these data were obtained at very low resolution
resulting in only three vectors coinciding with B1-E.
A strong magnetic field does not necessarily relegate turbulent compression or gravity to minor roles in the evolution of
B1-E. Indeed, Nakamura & Li (2008) suggest that core forming
regions exhibit different phases where magnetic fields, gravity, and turbulence are sometimes dominant and sometimes secondary. For example, they suggest that core forming regions begin with strong magnetic fields and strong turbulence, such that
gravity is a secondary effect. After the turbulent energy dissipates, gravitational collapse can proceed via magnetic field lines.
5. Conclusions
With recent Herschel dust continuum observations from the
Herschel Gould Belt Survey, we identified substructure in the
Perseus B1-E region for the first time. With our Herschel data,
we determined the temperature and column density across the region. We selected the nine highest column density substructures
for complementary observations with the GBT. We summarize
our main conclusions as follows:
1. B1-E contains a loose collection of roughly a dozen prominent substructures. This morphology is atypical of most star
forming regions, which produce dense clusters and organize
material along filaments. Such a loose collection of cores can
be produced in magnetically subcritical simulations (e.g., see
Li & Nakamura 2002; Basu et al. 2009a) influenced by a
strong magnetic field. Furthermore, a strong magnetic field
will delay the onset of star formation, which may explain the
age discrepancy between B1-E and the other nearby clumps
in Perseus, IC 348 and NGC 1333.
2. The B1-E clump as a whole is gravitationally bound with an
estimated virial parameter of α ≈ 0.9. Therefore, B1-E is
likely to form dense cores and stars. Assuming T = 14 K
and n = 5 × 103 cm−3 , we find a Jeans length within the entire B1-E clump of ∼0.35 pc, accurate within a factor of a
few. We measure a median nearest neighbour separation of
∼0.15 pc for our nine substructures (∼0.17 pc excluding B1E6 and B1-E7), for an inclination of 60◦ . Thus, the substructures have a median minimum separation that is less than the
Jeans length by a factor of two. This smaller length scale
could indicate that B1-E contracted after the observed substructures formed.
3. Several of the B1-E substructures have masses on order
of the Jeans mass (∼0.9 M ), assuming T = 10 K and
n = 1 × 105 cm−3 . Nevertheless, most substructures have
large virial parameters (α 2) indicating that they are gravitationally unbound. The large virial parameter suggests that
non-thermal motions are significant and that the Jeans mass
does not well represent the critical mass scale for these substructures. B1-E2 is the only substructure with a small virial
parameter (1 <
∼α<
∼ 3) and, due to its narrow line emission,
may be gravitationally bound. Furthermore, NH3 (2, 2) and
HC5 N (9–8) were only detected towards B1-E2. Thus, B1E2 is an excellent candidate for a core precursor.
4. The B1-E substructures have a substantial centroid velocity dispersion (∼0.42 km s−1 ) resulting in an interaction
timescale of ∼1 Myr, assuming an average separation of
∼0.46 pc. This timescale is well within the lifetime of protostellar cores. Thus, competitive accretion may play a significant role in the evolution of structure in B1-E.
S. I. Sadavoy et al.: Herschel observations of Perseus B1-E: a first look at core formation
The B1-E region appears to be an excellent candidate for future core formation. Further studies, however, are necessary
to explore the dust and dynamics of the bulk gas in B1-E.
For example, additional submillimetre continuum data along
the Rayleigh-Jeans tail (e.g., with SCUBA-2) will improve
SED-fitting and more accurately probe the dust emissivity index, β, size, and mass. To date, our Herschel observations of
B1-E are the only well detected continuum data for the region.
Additionally, our GBT observations are the only high resolution
kinematic data for B1-E, and we observed only the nine highest
column density substructures seen with Herschel. These data illustrate the potential of B1-E to be a core forming region. Future
observations of both the magnetic field strength and the turbulent
velocities are necessary to probe further how B1-E will evolve.
Acknowledgements. We thank the anonymous referee for comments that greatly
improved the discussion and narrative of this paper. This work was possible
with funding from the Natural Sciences and Engineering Research Council of
Canada CGS award. We thank B. Ali, S. Basu, R. Friesen, D. Johnstone, V.
Konyves, G. Langston, A. Pon, S. Schnee, B. Schulz, N. Wityk, and the kind
folks at the NHSC helpdesk and at the GBT for their invaluable assistance
and advice throughout this project. We also thank the CITA IT staff for their
time and attention. SPIRE has been developed by a consortium of institutes
led by Cardiff Univ. (UK) with Univ. Lethbridge (Canada); NAOC (China);
CEA, LAM (France); IFSI, Univ. Padua (Italy); IAC (Spain); Stockholm
Observatory (Sweden); Imperial College London, RAL, UCL-MSSL, UKATC,
Univ. Sussex (UK); Caltech, JPL, NHSC, Univ. Colorado (USA). This development has been supported by national funding agencies: CSA (Canada);
NAOC (China); CEA, CNES, CNRS (France); ASI (Italy); MCINN (Spain);
SNSB (Sweden); STFC (UK); and NASA (USA). PACS has been developed
by a consortium of institutes led by MPE (Germany) with UVIE (Austria); KU
Leuven, CSL, IMEC (Belgium); CEA, LAM (France); MPIA (Germany); INAFIFSI/OAA/OAP/OAT, LENS, SISSA (Italy); IAC (Spain). This development
has been supported by the funding agencies BMVIT (Austria), ESA-PRODEX
(Belgium), CEA/CNES (France), DLR (Germany), ASI/INAF (Italy), and
CICYT/MCYT (Spain). The Green Bank Telescope is operated by the National
Radio Astronomy Observatory. The National Radio Astronomy Observatory is
a facility of the National Science Foundation, operated under the cooperative
agreement by Associated Universities, Inc.
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